Stellar structure

Stars of different mass and age have varying internal structures. Stellar structure models describe the internal structure of a star in detail and make detailed predictions about the luminosity, the color and the future evolution of the star.

Energy transport

Different layers of the stars transport heat up and outwards in different ways, primarily convection and radiative transfer, but thermal conduction is important in white dwarfs.

Convection is the dominant mode of energy transport when the temperature gradient is steep enough so that a given parcel of gas within the star will continue to rise if it rises slightly via an adiabatic process. In this case, the rising parcel is buoyant and continues to rise if it is warmer than the surrounding gas; if the rising particle is cooler than the surrounding gas, it will fall back to its original height. [harvtxt|Hansen|Kawaler|Trimble|2004|loc=§5.1.1] In regions with a low temperature gradient and a low enough opacity to allow energy transport via radiation, radiation is the dominant mode of energy transport.

The internal structure of a main sequence star depends upon the mass of the star.

In solar mass stars (0.3–1.5 solar masses), including the Sun, hydrogen-to-helium fusion occurs primarily via proton-proton chains, which do not establish a steep temperature gradient. Thus, radiation dominates in the inner portion of solar mass stars. The outer portion of solar mass stars is cool enough that hydrogen is neutral and thus opaque to ultraviolet photons, so convection dominates. Therefore, solar mass stars have radiative cores with convective envelopes in the outer portion of the star.

In massive stars (greater than about 1.5 solar masses), the core temperature is above about $1.8\; imes\; 10^7$ K, so hydrogen-to-helium fusion occurs primarily via the CNO cycle. In the CNO cycle, the energy generation rate scales as the temperature to the 15th power, whereas the rate scales as the temperature to the 4th power in the proton-proton chains. [harvtxt |Hansen|Kawaler|Trimble|2004|loc=Tbl. 1.1] Due to the strong temperature sensitivity of the CNO cycle, the temperature gradient in the inner portion of the star is steep enough to make the core convective. In the outer portion of the star, the temperature gradient is shallower but the temperature is high enough that the hydrogen is nearly fully ionized, so the star remains transparent to ultraviolet radiation. Thus, massive stars have a radiative envelope.

The lowest mass main sequence stars have no radiation zone; the dominant energy transport mechanism throughout the star is convection. Giants are also fully convective. [harvtxt |Hansen|Kawaler|Trimble|2004|loc=§2.2.1]

Equations of stellar structure

The simplest commonly used model of stellar structure is the spherically symmetric quasi-static model, which assumes that a star is in a steady state and that it is spherically symmetric. It contains four basic first-order differential equations: two represent how matter and pressure vary with radius; two represent how temperature and luminosity vary with radius. [This discussion follows those of, e. g., harvtxt |Zeilik|Gregory|1998|loc=§16-1–16-2 and harvtxt |Hansen|Kawaler|Trimble|2004|loc=§7.1.]

In forming the stellar structure equations (exploiting the assumed spherical symmetry), one considers the matter density $ho(r)$, temperature $T(r)$, total pressure (matter plus radiation) $P(r)$, luminosity $l(r)$, and energy generation rate per unit mass $epsilon(r)$ in a spherical shell of a thickness $mbox\{d\}r$ at a distance $r$ from the center of the star. The star is assumed to be in local thermodynamic equilibrium (LTE) so the temperature is identical for matter and photons. Although LTE does not strictly hold because the temperature a given shell "sees" below itself is always hotter than the temperature above, this approximation is normally excellent because the photon mean free path, $lambda$, is much smaller than the length over which the temperature varies considerably, i. e. $lambda\; ll\; T/|\; abla\; T|$.

First is a statement of "hydrostatic equilibrium:" the outward force due to the pressure gradient within the star is exactly balanced by the inward force due to gravity.:$\{mbox\{d\}\; P\; over\; mbox\{d\}\; r\}\; =\; -\; \{\; G\; m\; ho\; over\; r^2\; \}$,where $m(r)$ is the cumulative mass inside the shell at $r$ and "G" is the gravitational constant. The cumulative mass increases with radius according to the "mass continuity equation:":$\{mbox\; \{d\}\; m\; over\; mbox\{d\}\; r\}\; =\; 4\; pi\; r^2\; ho\; .$

Integrating the mass continuity equation from the star center ($r=0$) to the radius of the star ($r=R$) yields the total mass of the star.

Considering the energy leaving the spherical shell yields the "energy equation:":$\{mbox\{d\}\; l\; over\; mbox\{d\}\; r\}\; =\; 4\; pi\; r^2\; ho\; (\; epsilon\; -\; epsilon\_\; u\; )$,where $epsilon\_\; u$ is the luminosity produced in the form of neutrinos (which usually escape the star without interacting with ordinary matter) per unit mass. Outside the core of the star, where nuclear reactions occur, no energy is generated, so the luminosity is constant.

The energy transport equation takes differing forms depending upon the mode of energy transport. For conductive luminosity transport (appropriate for a white dwarf), the energy equation is:$\{mbox\{d\}\; T\; over\; mbox\{d\}\; r\}\; =\; -\; \{1\; over\; k\}\; \{\; l\; over\; 4\; pi\; r^2\; \},$where "k" is the thermal conductivity.

In the case of radiative energy transport, appropriate for the inner portion of a solar mass main sequence star and the outer envelope of a massive main sequence star,:$\{mbox\{d\}\; T\; over\; mbox\{d\}\; r\}\; =\; -\; \{3\; kappa\; ho\; l\; over\; 64\; pi\; r^2\; sigma\; T^3\},$where $kappa$ is the opacity of the matter, $sigma$ is the Stefan-Boltzmann constant, and the Boltzmann constant is set to one.

The case of convective luminosity transport (appropriate for non-radiative portions of main sequence stars and all of giants and low mass stars) does not have a known rigorous mathematical formulation. Convective energy transport is usually modeled using mixing length theory. Mixing length theory treats the gas in the star as containing discrete elements which roughly retain the temperature, density, and pressure of their surroundings but move through the star as far as a characteristic length, called the "mixing length". [harvtxt|Hansen|Kawaler|Trimble|2004|loc=§5.1] For a monatomic ideal gas, mixing length theory yields:$\{mbox\{d\}\; T\; over\; mbox\{d\}\; r\}\; =\; left(1\; -\; \{1\; over\; gamma\}\; ight)\; \{T\; over\; P\; \}\; \{\; mbox\{d\}\; P\; over\; mbox\{d\}\; r\},$where $gamma\; =\; c\_p\; /\; c\_v$ is the adiabatic index, the ratio of specific heats in the gas. (For a fully ionized ideal gas, $gamma\; =\; 5/3$.)

Also required is the equation of state, relating the pressure to other local variables appropriate for the material, such as temperature, density, chemical composition, etc. Relevant equations of state may have to include the perfect gas law, radiation pressure, pressure due to degenerate electrons, etc.

Combined with a set of boundary conditions, a solution of these equations completely describes the behavior of the star. Typical boundary conditions set the values of the observable parameters appropriately at the surface ($r=R$) and center ($r=0$) of the star: $P(R)\; =\; 0$, meaning the pressure at the surface of the star is zero; $m(0)\; =\; 0$, there is no mass inside the center of the star, as required if the mass density remains finite; $m(R)\; =\; M$, the total mass of the star is the star's mass; and $T(R)\; =\; T\_\{eff\}$, the temperature at the surface is the effective temperature of the star.

Although nowadays stellar evolution models describes the main features of color magnitude diagrams, important improvements have to be made in order to remove uncertainties which are linked to our limited knowledge of transport phenomena. The most difficult challenge remains the numerical treatment of turbulence. Some research teams are developing simplified modelling of turbulence in 3D calculations.

ee also

*Polytrope

References

General references

*citation| title=Stellar Structure and Evolution | first1=R. | last1=Kippenhahn | first2=A. | last2=Weigert | publisher=Springer-Verlag | year=1990
*citation|last=Hansen | last2=Kawaler | last3=Trimble | first=Carl J. | first2=Steven D. | first3=Virginia | publisher=Springer | edition=2nd | year=2004 | title=Stellar Interiors | isbn=0387200894
*citation | last1=Kennedy | first1=Dallas C. | last2=Bludman | first2=Sidney A. | title=Variational Principles for Stellar Structure | year=1997 | journal=Astrophysical Journal | volume=484 | pages=329 | id=arxiv|astro-ph|9610099 | doi=10.1086/304333
*citation | first1=Achim | last1=Weiss | first2=Wolfgang | last2=Hillebrandt | first3=Hans-Christoph | last3=Thomas | first4=H. | last4=Ritter | title=Cox and Giuli's Principles of Stellar Structure | publisher=Cambridge Scientific Publishers | year=2004
*citation | last=Zeilik | first=Michael A. | last2=Gregory | first2=Stephan A. | title=Introductory Astronomy & Astrophysics | edition=4th | year=1998 | publisher=Saunders College Publishing | isbn=0030062284

External links

* [http://www-pat.llnl.gov/Research/OPAL OPAL opacity code]
* The [http://astro.ensc-rennes.fr/index.php?pw=ycesam Yellow CESAM code] , stellar evolution and structure FORTRAN source code
* [http://theory.kitp.ucsb.edu/%7Epaxton/EZ-intro.html EZ to Evolve ZAMS Stars] a FORTRAN 90 software derived from Eggleton's Stellar Evolution Code, a web-based interface can be found [http://shayol.bartol.udel.edu/~rhdt/ezweb here] .
* [http://obswww.unige.ch/~mowlavi/evol/stev_database.html Geneva Grids of Stellar Evolution Models] (some of them including rotational induced mixing)
* The [http://www.oa-teramo.inaf.it/BASTI BaSTI] database of stellar evolution tracks